Lecture 7 GAS in Galaxies Lec 1

Lecture 7 GAS in Galaxies Lec 1
How to Calculate Your Own Star
•  Today on astro-ph
•  arXiv:1509.06775 [pdf, ps, other] grayStar3 - gray no more: More physical realism and
a more intuitive interface - all still in a WWW browser C. Ian Short •  The goal of the openStar project is to turn any WWW browser, running on any platform,
into a virtual star equipped with parameter knobs and instrumented with output
displays that any user can experiment with using any device for which a browser is
available. .... The code integrates scientific modeling in JavaScript with output
visualization HTML. The user interface is adaptable so as to be appropriate for a large
range of audiences from the high-school to the introductory graduate level. The
modeling is physically based and all outputs are determined entirely and directly by the
results of in situ modeling, giving the code significant generality and credibility for
pedagogical applications.
•  gS3 also models and displays the circumstellar habitable zone (CHZ) and allows the
user to adjust the greenhouse effect and albedo of the planet. In its default mode the code
is guaranteed to return a result within a few second of wall-clock time on any device. The
more advanced user has the option of turning on more realistic physics modules that
address more advanced topics in stellar astrophysics. •  gS3 is a public domain, open source project and the code is available from
www.ap.smu.ca/~ishort/grayStar3/ and is on GitHub. gS3 effectively serves as a public
library of generic JavaScript+HTML plotting routines that may be recycled by the
community.
GAS
The other baryonic component- sec 2.4 in S+G Material scattered in Ch 8-9 of MWB
See web page of Alyssa Goodman at Harvard Astronomy 201b :
Interstellar Medium and Star Formation
http://ay201b.wordpress.com/
I will be going thru material a bit too fast for derivations and
strongly recommend looking at the above pages for details
See also Molecular Gas in Galaxies ARA&A Vol. 29: 5811991J. S. Young and N. Z.
Scoville Dopita, M., & Sutherland, R.: Astrophysics of the Diffuse Universe 2005 Lequeux, J.: The Interstellar Medium, Springer, Berlin, 2003
Osterbrock, D.E., & Ferland, G.J.: Astrophysics of Gaseous Nebulae and
Active Galactic Nuclei, Palgrave Macmillan, 2006
Spitzer, L.: Physical Processes in the Interstellar Medium1978 Thronson, H.A., Shull, J.M. (Herausgeber) : The Interstellar Medium in
Galaxies, Kluwer Academic Publishers, 1990 Gas- Big Picture •  Dark matter halos grow by merging and accretion (e.g.
Galaxies can grow by accretion of gas, by merging with gas
rich galaxies and by merging with gas poor galaxies)
•  Gas falls into these halos, cools and forms stars. •  How does this occur- the physics of gas accretion,
–  How and when did galaxies accrete their gas and what
do they do with it (e.g. form ISM, stars, expel the gas,
feed the supermassive black hole ....) Paper to Read
•  Please read http://arxiv.org/
abs/0906.0605 Dotter et al in
Proceedings of IAU Symposium
258: The Ages of Stars
•  How well can we measure the
ages of stars • 
cold gas: dominates in
Spirals-many phases –  neutral hydrogen
–  molecular gas-Dense
molecular clouds, have
most of the total mass of
the interstellar gas and
are of key importance for
star formation, occupy a
negligible fraction of the
total volume
–  warm ionized gas-has
persistent transient states
out of thermal pressure
balance •  Milky-Way-like galaxies
cold gas mass~10% of the
stars
•  For lower mass galaxies
the baryonic fraction in
gas is larger; at
Mstar<109.5M!
gas dominates the
baryonic content
•  Hot gas (T~106-7 k)
dominant ISM in elliptical
galaxies In spirals hot gas fills the
volume but low total mass GAS-ISM
x-ray images of elliptical galaxies
emphasing structure
GAS-ISM
x-ray images of elliptical galaxies
emphasing structure
log gas fraction for star forming galaxies
Gas
•  Other than stars the baryons in
galaxies lie in 3 forms
–  gas
–  rocks
–  dust (0.1% of mass) •  the % mass in rocks and dust is
small
•  There is an interplay between the
stars and gas,with stars forming out
of the gas and with enriched gas
being ejected back into the
interstellar medium from evolved
stars.
•  There exist a vast array of spectral
diagnostics for the gas in both
emission and absorption which can
reveal –  chemical composition
–  temperature
–  velocities
Peeples and Shankar 2011
–  ionization mechanism A Bit of Physics •  Saha equation describes the
ionization balance of the gas which
depends on the temperature,
quantum mechanical transition
probabilities and densities •  An atom with multiple energy states
in thermal equilibrium with a
radiation field will find itself in one or
another of these energy states. •  Frequent transitions to and from
other states will occur as photons
interact with the atoms.
•  transitions from the upper of the
states of figure take place by photo
deexcitation and by induced
deexcitation. Transition in the
upward direction is by
photoexcitation or collisional
excitation
For lots of details
see MBW appendix B
A Bit of Physics •  The rates of ionization and
recombination are important (see eqs
2.21,2..22 in S+G); e.g. X++e
X+γ
•  the rate at which ions recombine thus
clearly depends on the ion density, X+
and the electron density and the
recombination coefficient, α, which
depends on the ion, (e.g. the number
of electrons it has and its atomic
number) •  Thus recombination rate of electrons
for a given ion X++ is dne/dt=nx+neα(Te); •  the recombination time is the #of
electrons/ the rate :ne/dne/dt a few thousand years in a HII region α the recombination rate depends on
QM and Boltzmann's law In steady state # of
ionizations= # of
recombinations
Ionization is from
•  collisions with hot
electrons • photoionization from stars
• shocks
Atomic Lines •  The energy levels and
transitions for hydrogen •  Each element and ionization set
has a similar (but more
complex) set of lines
•  The probability of emitting a
given line depends on the
temperature and density of the
gas A Bit of Physics-Ionizing Photons •  One can estimate the number of ionizing photons from a star using the
black body formula (e.q. 1.35 in S&G) and integrating over the
photons more energetic than the ionization potential of the ion of
interest (e.g. H with13.6 eV)- effects of radiative transfer in stellar
atmospheres of hot stars is VERY important •  These photons ionize and heat the gas
•  The gas responds by emitting lines characteristic of the chemical
composition, temperature, ionization state, density etc ...
•  Please see https://ay201b.wordpress.com/2011/04/12/course-notes/
#the_sound_speed for a LOT more detail (also covered in radiative
processes course) Physics of Emission from Gas-MWB
sec 10.3.7 •  Gas is heated/excited/ionized by photons (stars, AGN), shocks
(supernova) and gravity •  Atomic transitions reveal the ionization state, temperature, density,
velocity structure and chemical composition of the gas. •  Photoionization: photon from source eject electron from ion- to do
this photon needs to have energy greater than ionization potential
(e.g. 13.6 eV for Hydrogen; O,B stars, AGN)
•  Collisional ionization: gas is excited by collisions with 'hot' electrons
(again electron energy has to be above threshold). Electrons have
Maxwell-Boltzman energy distribution in equilibrium •  wide range of types of transitions: 2 'basic' types
–  permitted: fast transition rate, line is emitted before ions state is
altered
–  forbidden: violate transition rule, ion can be collisionally deexcited when density exceeds critical density; presence of line
thus places constraint on gas density. - jargon forbidden lines are
indicated by [OII] (OII is the ionization state of the gas, once
ionized oxygen). Line Emission from Hydrogen (MBW
476-478)
•  Need detailed balance the flux F (number of photons per unit time)
has to be balanced by the recombination rate.
•  F=αΒNpNeV; αΒ is the recombination coefficient,Np is the proton
density,Ne is the electron density, V is volume.
•  If the region is optically thin the line emission corresponding to a
transition between states 1 and 2 is •  L12=4πε12V=hν12VNpNeα'
•  This gives for T=104K '
•  F=0.45hNpNeVνHα and Ηα/Ηβ=3.8'
• 
Thus, by measuring the luminosity of a HII region in a
recombination line, one can in principle infer the rate
which, in turn, can be used to infer the number of OB
stars that generate the ionizing photons'
A Bit of Physics-Relevant Velocities Sound speed in gas cs=∂P/∂ρ; P and ρ are the pressure and density For isothermal perfect gas P=ρkBT/µmH
cs=sqrt(kBT/µ)
where kB is Boltzmann's constant and µ is the mean molecular weight of
the gas Many astrophysical situations in the ISM are close to being isothermal,
thus the isothermal sound speed is often used
Reason: an increase in temperature due to compression will be
followed by radiative cooling to the original equilibrium temperature. Alfvén speed: The speed at which magnetic fluctuations propagate. vA = B /sqrt{4πρ} Alfvén waves are transverse waves along the
direction of the magnetic field. ISM- Relevant Velocities
Some characteristic values •  galactic rotation gradient 18km/sec/kpc
•  Thermal sound speed ideal gas for H, t is 0.3, 1, and 3 km/
s at 10 K, 100 K, and 1000 K- most of the velocites
measured are supersonic (e.g. gas is tubulent) •  Alfven speed- for typical ISM values B=1µG, # density n~1
cm-3 vA= 2 km/sec
A Bit of Physics-TimeScales In gas at temperature T, the mean particle velocity is given by the 3-d
kinetic energy: 3/2mv2 = kT; collision timescale τ~l/v; l ~ nσ; n is the NUMBER density of the gas and
σ is a typical cross section (hard sphere approx for ions πr2~10-15
cm-3 )
and thus τ~{2/3} {kTm}-1/2/(nσ) = 4.5 103n-1T-1/2 years
for a typical place in the ISM (n,T) = (1cm-3, 104)the collision time is 45
years
For a sphere of gas, where thermal pressure is balanced by self-gravity
the timescale to collapse (the Jeans time) is τJ~1/sqrt(4πGρ) which is
similar to the free falltime τff=(3π/32Gρ)1/2 = 4.4 x 104 yr /sqrt{nH /106}
https://en.wikipedia.org/wiki/Jeans_instability; nHis the particle density
ρ is the mass density
Simple Derivation of Jeans Collapse
•  Kinetic energy in cloud is KE=(3/2kT)N; N is the number of
particles, T is the temperature
•  The gravitational (binding) energy U=-3/5GM2/R (uniform
density sphere- derivation later in class) •  Using the viral theorm (lots more later) system is in equilibrium if 3NkT=(3/5)GM2/R
•  So to collapse the internal energy <binding energy •  Assume all the mass is in hydrogen with a mass m per
particle
•  then to collapse M>(5kT/Gm)3/2(3/4πρ)1/2 where ρ is the
density (e.g. (M/[4/3pπr3])
•  M is called the Jeans mass •  The material to read is Spitzer 1978- Physical Processes in
the ISM pg 286-287 (see web page for text) Big Questions •  What is the volume filling factor
of the hot ISM?
•  What is the distribution of the
temperature, density, and
velocity •  What are typical scales in the
ISM and why?
•  What is the effect of turbulence ,
magnetic fields and cosmic
rays? •  What causes density and
pressure inhomogeneities in the
evolution of the ISM?
•  How is the ISM related to star
formation? •  Why is the ISM in spirals and
ellipticals so different in density
and temperature?
x-ray temperature
map of LMC
intensity
map of LMC in x-rays
Physics of Emission from Gas •  Lines have enormous range of energies/wavelengths
–  molecular and fine structure lines in IR/radio band
–  atomic lines in the IR, optical, UV and x-ray •  Ionized gas also emits a continuum via thermal bremmstrahlungshape of which is a measure of temperature, intensity goes as density
squared (board) •  Observed line energies give velocity information: redshift, velocity
field
•  Relative strength of lines determines ionization temperature,
abundance of given element (corrected for ionization balance (go to
board)). •  see Thermal radiation processes J.S. Kaastra, F.B.S. Paerels,
F. Durret, S. Schindler, P. Richter
Space Science Reviews, Volume 134, Issue 1-4, pp. 155-190, 2008
astro-ph/0801.1011 for the background physics –  Importance of the ISM
•  Despite its low mass, the ISM is
very important
•  crucial role in the star-gas cycle
in spirals and irregulars,
–  it facilitates ongoing (&
current) star formation
–  it is a repository for elements
created in SNR and stars
and therefore is a key to
measure chemical evolution
•  Because it can cool, its collapse
is dissipational
–  stars can form !! hot gas
cold gas
stars:
–  galaxies are smaller than
dark matter halos !
its emission & absorption provides enormous diagnostic information
• Doppler motions reveal galaxy
dynamics
• Abundance measurements allow
study of chemical evolution
• physical conditions: density; temp;
pressure; turbulence; gas column
density; mass, • can all be derived from
observations of emission/
absorption lines
• lines are bright and can be
seen (relatively) easily at
cosmological distances.
The ISM in Spirals is DYNAMIC
• 
• 
There is strong
interaction
between the
different phases of
the ISM and
feedback between
star formation and
the rest of the ISM
There is lots of
complex nonlinear effects (and
lots of jargon)
Fabian Walter At low redshift ISM in spirals not affected much by AGN
Its not so clear if ISM in ellipticals is dynamic in the
same way; AGN seem to be more important How Does One Observe the ISM
(sec 5.2 in S&G) •  Because of the wide range in temperatures and
densities a wide variety of techniques are needed
•  Radio:
–  free-free emission and 21cm for HI
–  high freq radio-far IR (CARMA, ALMA,
Herschel) wide variety of molecular lines
•  IR spectral lines [OI]63,145um and [CII]158
and [CI]370,609um
•  Optical/UV
–  wide variety of emission and absorption
lines from ionized metals (C,N,O etc) - gas is
photoionized
•  Soft x-ray
•  continuum and emission lines from
T~106-107k gas (spirals and ellipticals)
- gas is collisionally ionized
γ-ray
•  interaction of cosmic rays with gas Hα Emission from MW
LMC in soft x-rays
Far IR Lines •  More than 145 lines , most of them rotationally excited lines from abundant molecules:
•  38 12CO lines (up to J=42-41 37 lines of both o-H2O and p-H2O (up to 818-717),16 OH
lines 12 13CO lines (up to J=16-15) and several HCN and HCO+ lines Goicoechea et al
2015 ApJ 799 102 ; brightest line is [OI] at 63µ •  this paper (Goicoechea et
al 2015 ApJ 799 102 ) is
very recent data from the
Herschel observatory on
Far IR lines from a star
forming region but is too
detailed for a discussion in
the class. •  Spiral ISM 'States'- f is the filling factor
•  Molecular Medium (MM): T~20 K, n > 103 cm-3, f < 1%. The MM is
mostly cold dense molecular clouds which are gravitationally bound.
this phase contains ~as much mass as the atomic hydrogen, but
occupies only a very small fraction of the ISM.
•  Cold Neutral Medium (CNM; T~100 K, n~20 cm-3, f =2 - 4%). The CNM
is distributed in rather dense filaments or sheets, occupying a minor
fraction of the ISM. The CNM is most readily traced by HI measured in
absorption.
•  - Warm Neutral Medium (WNM; T~6000 K, n~ 0.3 cm-3, f~30%). This
phase provides the bulk of the HI seen in emission line surveys.
•  - Warm Ionized Medium (WIM; T ~8000 K, n ~0.3 cm-3 f~15%).
associated with HII regions, but a considerable fraction of the ISM
outside of HII regions is also filled with ionized gas.
•  - Hot Ionized Medium (HIM; T~106 K, n~10-3 cm-3 , f~50%). The hot gas
produced by supernova explosions and their after effects (in spirals,
other physics in ellipticals( - long cooling time , a large fraction of the
ISM is filled with this component.
•  http://ned.ipac.caltech.edu/level5/March01/Brinks/Brinks4.html
ISM- Phases
•  Hot ionized medium (e.g. X-rays)
•  Warm ionized medium HII region(e.g. Hα)
•  Warm neutral medium (e.g. HI emission)
•  Cold neutral medium (e.g. HI absorption)
•  Molecular medium (e.g. CO)
These phases have different distributions perpendicular to the planescale height Fabian Walter
The ISM
• 
• 
• 
• 
The 5 'states' are in dynamic interaction. !
the coldest clouds are molecular and the densest (hydrogen molecules, CO, NH3 and other
molecules)- this is where stars form . !
The dust is composed of 'refractory' elements and molecules mainly carbon, silicon, iron
and is responsible for most of the absorption of optical light in the galactic plane - the
energy absorbed by the dust heats it and the dust re-radiates in the IR !
The ISM is threaded by magnetic fields. At ~ 5µG, these fields provide a pressure
comparable to the pressure of the gas . The magnetic fields therefore affects the dynamics
of the ISM!
•  Book on the subject Bruce Draine ' Physics of the
Interstellar and Intergalactic Medium' Princeton
series on Astrophysics !
Optical spectrum of HII Region •  Optical spectrum show lines due to [OII]. [OIII],Hα, [NII], etc Molecular Lines
•  Molecular clouds are very rich in spectral features from a wide
variety of molecules- lots of information
•  Some of the lines (CO) are so strong that they can be seen at high
redshift
Millimeter Band Spectrum of Molecular
Cloud Millimeter Band Spectrum of Molecular
Cloud •  zoom in on previous plot How do Molecules Emit Radiation •  Emission is primarily
from rotational and
vibrational levels –  Millimeter emission:
rotational transitions
Infrared absorption:
vibrational transitions
Limitation: need
background IR source =>
only info along line of
sight
•  Earth s atmosphere
prevents observations of
key molecules from
ground: H2O, O2, CO2
Ewine F. van Dishoeck
MM emission:Limitation: molecule must
have permanent dipole moment => cannot
observe H2, C2, N2, CH4, C2H2, …
Advantage: many molecules down to low
abundances;
Gas Cooling
•  Collisional excitation: free electron impact knocks a bound electron to
an excited state; it decays,emitting a photon.
•  Collisional ionization: free electron impact ionizes a formerly bound
electron, taking energy from the free electron.
•  Recombination: free electron recombines with an ion; the binding
energy and the free electron's kinetic energy are radiated away •  Free-free emission: free electron is accelerated by an ion, emitting a
photon. (A.k.a. Bremsstrahlung)
Reference on Cooling via molecular rotational lines and dust emission
(Neufeld, Lepp and Melnick (1995, Ap.J.Supp., 100, 132) Gas Cooling L=n2Λ(Τ)
MWB sec 8.1.3, 8.4
•  T>107k thermal
bremmstrahlung L~n2T 1/2
•  107>kT>106.3k Fe L lines
•  104.5>kT>106.3k K and L
lines of 'metals'
•  104>kT>104.5k Hydrogen •  At lower temperatures fine
structure lines and
molecules dominate log T
Cooling curve as a function
of kT and metallicity-for gas
in collisional equilibrium Sutherland and Dopita
table 2.5 in S&G
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